AbstractAims. We present a detailed abundance analysis of a
strongly r-process enhanced giant star discovered
in the HERES project, HE 2327-5642, for which [Fe/H] = -2.78,
[r/Fe] = +0.99. Methods. We determined the stellar parameters and
element abundances by analyzing the high-quality VLT/UVES spectra. The
surface gravity was calculated from the non-local thermodynamic
equilibrium (NLTE) ionization balance between Fe ,i and Fe II,
and Ca I and Ca II.
Results. Accurate abundances for a total of 40
elements and for 23 neutron-capture elements beyond Sr and up to Th
were determined in HE 2327-5642. For every chemical species,
the dispersion in the single line measurements around the mean does not
exceed 0.11 dex. The heavy element abundance pattern of
HE 2327-5642 is in excellent agreement with those previously
derived for other strongly r-process enhanced
stars, such as CS 22892-052, CS 31082-001, and
HE 1219-0312. Elements in the range from Ba to Hf match the
scaled Solar r-process pattern very well. No firm
conclusion can be drawn about the relationship between the fisrt
neutron-capture peak elements, Sr to Pd, in HE 2327-5642 and
the Solar r-process, due to the uncertainty in the
Solar r-process. A clear distinction in Sr/Eu
abundance ratios was found between the halo stars of different europium
enhancement. The strongly r-process enhanced stars
contain a low Sr/Eu abundance ratio at [Sr/Eu
,
while the stars with 0 < [Eu/Fe] < 1 and [Eu/Fe] <
0 have 0.36 dex and 0.93 dex higher Sr/Eu values,
respectively. Radioactive dating for HE 2327-5642 with the
observed thorium and rare-earth element abundance pairs results in an
average age of 13.3 Gyr, when based on the high-entropy wind
calculations, and 5.9 Gyr, when using the Solar r-residuals.
We propose that HE 2327-5642 is a radial-velocity variable
based on our high-resolution spectra covering 4.3 years.

1 Introduction

The detailed chemical abundances of Galactic halo stars contain unique
information about the history and nature of nucleosynthesis in our
Galaxy. A
number of observational and theoretical studies have established that
in the
early Galaxy the rapid (r) process of neutron
captures was primarily
responsible for the formation of heavy elements beyond the iron group
(we
cite only the pioneering papers of Spite & Spite 1978; Truran 1981).
The onset of the slow
(s) process of neutron captures occurred at later
Galactic times (and higher
metallicities) with the injection of nucleosynthetic material from
long-lived
low- and intermediate-mass stars into the interstellar medium (see Travaglio et al. 1999,
and
references therein). Since 1994, a few rare stars have been found that
exhibit large enhancements of the r-process
elements, compared to Solar
ratios, suggesting that their observed abundances are dominated by the
influence of a single, or at most very few nucleosynthesis events. The
r-process is associated with explosive conditions of
massive-star
core-collapse supernovae (Woosley
et al. 1994), although the astrophysical
site(s) of the r-process has yet to be identified.
Observations of stars with strongly enhanced r-process
elements have placed important
constraints on the astrophysical site(s) of their synthesis.

Sneden et al. (1994)
found that the extremely metal-poor
([Fe/H] -3.1)
giant CS 22892-052 is neutron-capture-rich, [Eu/Fe]
(following
the suggestion of Beers &
Christlieb 2005, we hereafter refer to stars having
and
as
r-II stars), and that
the relative abundances of nine elements in the range from Ba to Dy are
consistent with a scaled Solar System r-process
abundance distribution.
Later studies of CS 31082-001 (Hill
et al. 2002), BD+17 3248
(Cowan et al. 2002),
CS 22892-052 (Sneden
et al. 2003), HD 221170 (Ivans et al. 2006),
CS 22953-003 (François
et al. 2007), HE 1219-0312, and
CS 29491-069
(Hayek et al. 2009)
provided strong evidence of a universal production ratio of
the second r-process peak elements from Ba to Hf
during the Galaxy history.
CS 31082-001 (Hill
et al. 2002) provided the first solid evidence that
variations in progenitor mass, explosion energy, or other intrinsic and
environmental factors or all of these may produce significantly
different r-process yields
in the actinide region ().
The third r-process peak (
)
has not been well constrained, because, in most r-II stars, it is only
probed by abundance
measurements of two elements, Os and Ir. The abundances of platinum and
gold were obtained for CS 22892-052 (Sneden
et al. 2003). The only detection of lead in a r-II
star so far is in CS31082-001 (Plez
et al. 2004).

Sneden et al. (2003)
reported an underabundance of elements in the range of
40 < Z <
56 relative to the scaled Solar r-process, which
prompted a discussion
of multiple r-process sites (see, for example, Travaglio
et al. 2004; Farouqi et al. 2009; Qian &
Wasserburg 2008). The detection of the radioactive elements
thorium and uranium provided new
opportunities for deriving the ages of the oldest stars and hence
determining a
lower limit to the age of the Universe (see the pioneering papers
of Cayrel
et al. 2001; Sneden et al. 1996). It
appears that all the r-II stars with measured Th (and U) can
be divided into two groups: (a) stars exhibiting an actinide
boost (e.g., CS 31082-001, HE 1219-0312), and
(b) stars with no obvious
enhancement of thorium with respect to the scaled Solar r-process
pattern
(e.g., CS 22892-052, CS 29497-004; for a full list of
stars, see
Roederer et al. 2009).
For the actinide boost stars, ages cannot be derived
when only a single radioactive element, either Th or U, is
detected.

Continuing our series of papers on the Hamburg/ESO R-process-Enhanced
Star
survey (HERES), we aim to extend our knowledge of heavy
element synthesis in the early Galaxy by means of a detailed abundance
analysis of the
strongly r-process enhanced star
HE 2327-5642. We also
investigate the reliability of multiple Th/X
chronometers for
HE 2327-5642, where X is an element in the
Ba-Hf range.

HE 2327-5642 was identified as a candidate metal-poor
star in the Hamburg/ESO Survey (HES; see Christlieb et al.
(2008) for
details of the candidate selection procedures). Moderate-resolution
(
Å) spectroscopy
obtained at the Siding Spring
Observatory (SSO) 2.3 m-telescope with the Double Beam
Spectrograph (DBS)
confirmed its metal-poor nature. Therefore, it was included in the
target list
of the HERES project. A detailed description of the project and its
aims can
be found in Christlieb
et al. (2004, hereafter Paper I), and the
methods of automated
abundance analysis of high-resolution ``snapshot'' spectra were
described
in Barklem
et al. (2005, hereafter Paper II).
``Snapshot'' spectra with a spectral resolution
and
a signal-to-noise ratio
per pixel at 4100 Å were used to show that
HE 2327-5642 exhibits strong overabundances of the r-process
elements,
with [Eu/Fe] = +1.22 and
(Paper II).

This paper is structured as follows. After describing the
observations in Sect. 2, we
describe our abundance analysis of
HE 2327-5642 in Sects. 3
and
4, based
on high-quality VLT/UVES spectra and MAFAGS model
atmospheres (Fuhrmann et al. 1997).
The heavy element abundance pattern of HE 2327-5642 is
discussed in
Sect. 5.
Section 6
reports on the radioactive decay age determination. Our conclusions are
presented in Sect. 7.

2 Observations

For the convenience of the reader, we list the coordinates and
photometry of HE 2327-5642 in Table 1. The
photometry was taken from Beers
et al. (2007). High-quality spectra of
this star was acquired during May-November 2005 with the VLT and UVES
in
dichroic mode. The BLUE390+RED580 (4 h total integration time)
and
BLUE437+RED860 (10 h) standard settings were employed to
ensure a wide
wavelength coverage. The slit width of 0.8'' in both arms yielded a
resolving power of R=60 000. A pixel
binning ensured a proper
sampling of the spectra. The observations are summarized in
Table 2.

The pipeline-reduced spectra were shifted to the stellar rest
frame and then coadded
in an iterative procedure in which we identified pixels in the
individual spectra affected
by cosmic ray hits that had not been fully removed during the data
reduction, or those affected by
CCD defects or other artifacts. These pixels were flagged
and ignored in the final iteration of the coaddition. Both sets of
coadded blue
spectra have S/N of at least 50
per pixel at
Å. At the shortest
wavelengths, the S/N of the
BLUE390 and BLUE437 is 10 (at 3330 Å) and 70 (at
3756 Å),
respectively. The red arm spectra have S/N
> 100 per pixel in most of the
covered spectral range.

Barycentric radial velocities of HE 2327-5642 as
measured with Gaussian fits of selected absorption lines in our
high-resolution spectra, covering 4.3 years, indicate that the star is a
radial-velocity variable, although no signatures of a double-lined
spectroscopic binary star have been found. Our analysis of the data
taken during the
Modified Julian Date (MJD) period
53587.2-53666.2 (Table 3)
infers that the
radial velocity varies on timescales of 10 days, and that the radial
velocity curve underwent a minimum approximately at MJD 53620 (see
Fig. 1).
Furthermore, the measurement at MJD 54407.019 deviates
by 20 km s-1
from the average of the radial velocities measured at the
other epochs, and by a similar amount from the measurement taken only
about
three months later. The available data cannot be fitted satisfactorily
by a
sinusoidal curve, and we therefore suspect that the orbit of the system
is
highly elliptical. Additional observations are needed to confirm the
variability, and to determine both the period and the nature of the
orbit.

3 Analysis method

Our determinations of the stellar parameters and the elemental
abundances are
based on line profile and equivalent width analyses. We ignored any
lines with equivalent widths larger than 100 mÅ. Exceptions
were the elements, such as strontium, for which only strong lines can
be detected in HE 2327-5642. For a number of chemical
species, namely, H I, Na I,
Mg I, Al I,
Ca I-II, and Fe I-II,
non-local thermodynamic equilibrium
(NLTE) line formation was considered. The theoretical spectra of the
remaining
elements were calculated by assuming LTE. The coupled radiative
transfer and
statistical equilibrium equations were solved with the code NONLTE3
(Sakhibullin
1983; Kamp
et al. 2003) for H I and
Na I, and an
updated version of the DETAIL code (Butler
& Giddings 1985) for the remaining NLTE
species. The departure coefficients were then used to calculate the
synthetic
line profiles with the code SIU (Reetz
1991). The metal linelist was extracted from the VALD
database (Kupka et al. 1999).
For molecular lines, we
applied the data compiled by Kurucz
(1994). To compare with
observations, computed synthetic profiles were convolved with a profile
that combines both the instrumental broadening with a Gaussian profile
of 3.6 km s-1 and
the broadening caused by macroturbulence. We employ the macroturbulence
parameter
in the radial-tangential form as prescribed in Gray
(1992). By analyzing many line profiles in the spectrum of
HE 2327-5642, km s-1
was empirically found with some allowance to vary by 0.3 km s-1
(1).

The abundance analysis based on equivalent widths was
performed with the code
WIDTH9 (Kurucz
2004).
The SIU and WIDTH9 codes both treat continuum scattering correctly,
i.e.,
scattering is taken into account not only in the absorption
coefficient, but
also in the source function.

In both SIU and WIDTH9, we used the updated partition
functions from the
latest release of the MOOG code
(Sneden 1973), apart from those
of
Ho II and Ir II.
For Ho II, we adopted the partition
function calculated by Bord &
Cowley (2002). The Ir II
partition function was
revised based on the measured energy levels of van Kleef
& Metsch (1978). For the
temperature range with which we are concerned, this produces a
difference of +0.08/+0.2 dex in the Ho/Ir abundance
determined from the
Ho II/Ir I
lines.

3.1 Stellar parameters and atmospheric models

In Paper II, an effective temperature of
K
was derived from
photometry when adopting the reddening derived from the maps of
Schlegel et al. (1998).
A subsequent analysis of the snapshot spectrum inferred that
and
[Fe/H
.
For the stellar parameter
determination and abundance analysis, we used plane-parallel, LTE,
and line-blanketed MAFAGS model atmospheres (Fuhrmann
et al. 1997). Enhancements of the
-elements
Mg, Si, and Ca by the amounts determined in a close-to-final
iteration of our analysis were taken into account when computing these
model
atmospheres. Since suitable lines of oxygen were not covered by our
spectra in hand, we
could not determine the oxygen abundance, hence we adopted
,
which is typical of other stars of the same metallicity as
HE 2327-5642. We note that oxygen in cool
stellar atmospheres plays a minor role as both a donator of free
electrons and an opacity source, hence an uncertainty in the oxygen
abundance does not significantly
affect the calculated atmospheric structure.

Heiter & Eriksson
(2006) investigated the effect of geometry on atmospheric
structure and line formation for Solar abundance models, and concluded
that
plane-parallel models can be applied in abundance analyses for stars
with
and
K.
Therefore, HE 2327-5642 lies in the
stellar parameter range where the usage of plane-parallel models is
appropriate. This is confirmed by flux and abundance comparisons
between a
MAFAGS plane-parallel and MARCS (Gustafsson
et al. 2008)
spherical models with stellar parameters close to
those of HE 2327-5642, i.e.,
-3.
Synthetic spectra were computed for the wavelength range
3500-16 000 Å
with the code SIU, which solves the equation of radiative transfer in
only one
depth variable. For the absolute flux, we compared three different
combinations of model atmosphere and spectrum synthesis geometries
i.e.,
consistently plane-parallel (MAFAGS + SIU, p-p),
inconsistent (MARCS + SIU,
s-p), and consistently spherical
(MARCS model atmosphere library, s-s).
The
difference in absolute flux between these three models does not exceed
0.001 dex for wavelengths longer than 6600 Å. For
Å,
the p-p and s-p
fluxes are lower than those inferred from the s-smodel
with a maximum difference of 0.01 dex and 0.02 dex,
respectively, at
wavelengths around 3500 Å.

In Fig. 2,
we show the line profiles of H
and Hfor all
three models. The H
profile of the MAFAGS model is
consistent with that of the s-s
model. The difference in Hrelative
fluxes between the MAFAGS and s-s
models translates into an effective
temperature difference of 60 K. The abundance differences for
the selected
spectral lines were obtained between the p-p
and s-p models by fitting the
calculated synthetic spectra to the observed ones. The difference in
absolute
abundances,
,
is always negative, but does not exceed 0.01 and
0.02 dex for the lines of neutral and ionized species,
respectively.
The differences in abundance ratios are also negligible; i.e., smaller
than
0.01 dex.

Figure 2:

Synthetic profiles of H
( top panel) and H
( bottom panel) from the s-s
(dashed curve), s-p (continuous
curve), and p-p (dotted curve)
models. The calculations for the s-p
and p-p models were made for
pure hydrogen lines.

Top panel: synthetic flux profile of H computed
for K
(continuous curve) compared to the observed spectrum of
HE 2327-5642 (bold dots). The dashed curves show the effect of
a 80 K variation in the effective temperature on the synthetic
spectrum. In all calculations, we assumed
,
,
and km s-1.
Bottom panel: effective temperature derived from the H
(filled circles) and H
(open diamonds) line wings in HE 2327-5642 as a function of
surface gravity. The error bars show the uncertainty of
arising
from profile fitting.

The effective temperature of HE 2327-5642 was also determined
from a profile
analysis of H
and H
based on NLTE line formation calculations
of H I using the method described by Mashonkina et al.
(2008). Only these two
lines were employed because an accurate continuum rectification was not
possible in the spectral regions covering other Balmer lines. The
metallicity
and microturbulence velocity from Paper II were adopted during
the analysis
of these Balmer lines, while the gravity was varied between
and 2.4.
The theoretical profiles of H
and H
were computed by
convolving the profiles resulting from the thermal, natural, and Stark
broadening (Vidal
et al. 1970,1973), as well as self-broadening.
For the latter, we
use the self-broadening formalism of Barklem
et al. (2000).

We found that NLTE has a weak effect on the H
profile beyond
the core, because the difference between
derived
for this line assuming either NLTE or LTE does not exceed
20 K. We also found that the Hline wings are insensitive to
a variation in surface gravity across the stellar
parameter range with which we are concerned. The best fit solution was
achieved at
K.
Figure 3
(top panel) illustrates the quality of the
fits.

Based on S/N of the
observed spectrum and the sensitivity of the
Balmer lines to variations on
,
we estimate the uncertainty in
arising
from profile fitting to be 50 K for each line. For H,
NLTE leads to a weakening of the core-to-wing transition relative to
the LTE
case, resulting in a
that
is 80-100 K higher depending on surface gravity.
The effective temperature of HE 2327-5642 inferred from H
also
depends on ,
as shown in the bottom panel of Fig. 3. The
temperature deduced by combining the analyses of H
and H
is
K,
and a favorable range of
is that between 1.95 and 2.40.

For Ca and Fe, we applied a line-by-line differential NLTE
approach, in the
sense that stellar line abundances were compared with individual
abundances of
their Solar counterparts. With the adopted atomic
parameters, we note that the absolute Solar NLTE abundances obtained
from the two
ionization stages, Ca I and Ca II,
Fe I and Fe II,
were consistent within the error bars:
(Ca
I,
(Ca
II 8498 Å) = 6.29,
(Fe
I,
and (Fe
II(we refer to abundances on the
usual scale, where
).

We performed NLTE computations for a small grid of model
atmospheres with
two effective temperatures, namely
K,
derived from
photometry, and
K,
which is close to the result
of the Balmer line analysis. In the statistical equilibrium
calculations,
inelastic collisions with hydrogen atoms were accounted for using the
Steenbock & Holweger (1984)
formula with a scaling factor of
for
Ca and
for
Fe, as recommended by Mashonkina et al.
(2007a,2010).
The NLTE calculations for Ca I-II and Fe I-II
were iterated for various
elemental abundances until agreement between the theoretical and
observed
spectra was reached. The gravity was varied between
and 2.6 in steps of
0.2 dex. Microturbulence values were
tested in the range between
and 2.1 km s-1 in steps of
0.1 km s-1.
It was found that including the results for ,
2.0, and 2.1 km s-1 produced a
steep trend with measurable equivalent widths for the abundances found
from individual Fe I lines,
independent of the adopted values of
and
,
therefore these
values were excluded.

Adopting
K,
we obtained consistent iron abundances for the two
ionization stages if ,
2.32, and 2.32, and
values of 1.7,
1.8, and 1.9 km s-1,
respectively. For Ca, this is achieved for
,
2.37, and 2.47. Figure 4
illustrates the determination of
the surface gravity from the ionization equilibrium of Fe I/II
and Ca I/II when the remaining
stellar parameters are fixed at
K,
,
and km s-1.
When adopting
K,
the difference in
obtained from Fe and Ca does
not exceed 0.1 dex if
1.7 km s-1. Thus, we
identified two possible
combinations of stellar parameters for HE 2327-5642: (a)
5050/2.34/-2.78with km s-1,
and (b) 4980/2.23/-2.85 with km s-1
(see Fig. 5
for the combination
5050/2.34/-2.78). Both sets of the obtained parameters are
consistent with each other within the uncertainties in the stellar
parameters.

For consistency reasons, we adopted the effective temperature
adopted in
Paper II, i.e.,
K,
and the other stellar parameters determined in this study, i.e.,
,
,
and km s-1
(Table 4).
For the derived effective temperature and surface gravity, the
spectroscopic distance of HE 2327-5642 was estimated to range
from 4.4 to 4.9 kpc for stellar mass of
between 0.8 and 1 solar mass.

Figure 4:

NLTE abundances of Fe I (filled circles),
Fe II (open circles), Ca I
(filled diamonds), and Ca II (open
diamonds) in HE 2327-5642 as a function of surface gravity.
For clearer illustration, the symbols for Ca are shifted upwards by
0.5 dex. The calculations are for
K,
,
and km s-1.

Trends of abundances with equivalent width and excitation potential, as
determined from individual Fe I (filled
circles) and Fe II (open circles) lines,
using our adopted stellar parameters. The dotted line indicates the
mean Fe abundance from two ionization stages and the shaded grey area
its statistical error.

3.2 Line
selection and atomic data

The lines used in the abundance analysis were selected from the lists
of Paper II, Jonsell
et al. (2006), Lawler et al.
(2001c,2004),
Sneden et al. (2009),
and Ivans et al. (2006).
For atomic lines, we endeavored to apply single-source and recent gf-values
wherever possible, to diminish the uncertainties involved by combining
studies that may not
be on the same gf-value system. For the selected
lines of Na I,
Mg I, Al I, Ca I-II,
Sr II, and Ba II,
we
adopted gf-values (mostly from laboratory
measurements) and van der Waals
damping constants, which were carefully inspected in our previous
analyses of
the Solar spectrum (see Mashonkina
et al. 2008 for references).

Molecular data for two species, CH and NH, were assembled for
the abundance
determinations of carbon and nitrogen. For the analysis of the the A-X bands
at 4310-4313 Å and 4362-4367 Å, we used the CH line
list of Paper II, and we use the 13CH
line list described in Hill
et al. (2002).
The NH molecular line data for the A-X
band at 3358-3361 Å was taken
from Kurucz (1993).

The van der Waals damping for atomic lines was computed
following the ABO theory, where the data were
available, using the van der Waals damping
constants
at 10 000 K as provided by the VALD database (Kupka et al. 1999). We note
that the correct temperature dependence of the ABO theory was taken
into account. An exception was the selected lines of some elements, for
which we used the C6-values
derived from solar line-profile fitting by Gehren et al. (2004,
Na I, Mg I,
and Al I) and Mashonkina et al.
(2008, Sr II
and Ba II). If no other data were
available, the Kurucz
& Bell (1995)values
were employed.

Many elements considered here are represented by either a
single
isotope with an odd number of nucleons (Sc, Mn, Co, Pr, Tb, Ho, and Tm;
139La accounts for 99.9% of lanthanum according
to Lodders (2003)), or
multiple isotopes with measured wavelength differences (
Å
for Ca II, Ba II,
Nd II, Sm II,
Eu II, Yb II,
Ir I). Nucleon-electron spin interactions
in odd-A isotopes lead to hyper-fine splitting (HFS)
of the energy levels,
resulting in absorption lines divided into multiple components. Without
accounting properly for HFS and/or isotopic splitting (IS) structure,
abundances determined from the lines sensitive to these effects can be
severely overestimated. For example, in HE 2327-5642,
including HFS makes
a difference of -0.49 dex in the Ba abundance derived from the
Ba II
4554 Å line, and including IS leads to a
0.13 dex lower Ca abundance for
Ca II 8498 Å.

4 Abundance results

We derived the abundances of 40 elements from Li to Th in
HE 2327-5642,
and for four elements among them (Ca, Ti, Mn, and Fe), from two
ionization
stages. In Table 8
(online material), we list the results
obtained from individual lines. For every feature, we provide the LTE
abundance obtained and, for selected species, also the NLTE abundance.
In
Table 5,
we list the mean abundances, the dispersion in the single line
measurements about the mean (
),
and the
number of lines used to determine the mean abundances. We also list the
Solar photosphere abundances,
,
adopted from Lodders
et al. (2009), and the
abundances relative to iron, [X/Fe]. For the computation of [X/Fe],
[Fe/H]
was chosen as the reference,
with the exception of
the neutral species calculated based on a LTE assumption, where the
reference
is [Fe I/H]
.
We comment below on individual groups of
elements. The sample of cool giants from Cayrel
et al. (2004) was chosen as our comparison sample.

4.1 Li and CNO

With an equivalent width of 15 mÅ, the Li I
6708 Å line is
easily detected in this star. The abundance was determined using the
spectrum
synthesis approach, to account for the multiple-component structure of
the
line caused by both the fine structure of the upper energy level and
the presence of two isotopes, 7Li and 6Li.
The calculations of the
synthetic spectra were performed in two ways: (a) without 6Li,
and (b) by adopting the Solar isotopic ratio, i.e., 7Li:
6Li = 92.4:7.6
(Lodders 2003). In both
cases, the result for the Li abundance was
.
A goodness-of-fit analysis detected an
asymmetry in the Li I 6708 Å
line, which could be attributed to a
weak 6Li feature in the red wing of the 7Li
line. Although this asymmetry may also be convection-related (Cayrel et al. 2009), we
cannot exclude there being a significant amount of 6Li
in HE 2327-5642.
The departures from LTE cause only a
minor increase in the derived lithium abundance, i.e, by
0.04 dex according
to the calculations of Lind
et al. (2009).

With an abundance of
(Li)
= 0.99, HE 2327-5642 is located well
below the lithium plateau for halo stars near the main-sequence
turnoff, as
expected for a red giant (Iben 1967).

Figure 6:

Best fits (continuous curve) of the CH features near 4310 Å (
top panel) and 4211 Å ( middle panel),
and the NH molecular band near 3360 Å ( bottom panel).
The observed spectrum of HE 2327-5642 is shown as bold dots.
The dashed curves in the top and bottom panels
show the synthetic spectra with no carbon and nitrogen in the
atmosphere. In the middle panel, the continuous
curve corresponds to an isotope ratio of
,
while the dashed curves are synthetic spectra for
and 30.

Carbon was measured using CH lines in the regions 4310-4314 Å
and
4362-4367 Å, which are almost free from intervening atomic
lines (see
Fig. 6,
top panel). The C abundances obtained from these
spectral bands are consistent with each other to within
0.03 dex (see
Table 8,
online material). The mean abundance is
,
which is similar to those of the giants with
K
from the sample of Cayrel
et al. (2004).

We were able to use the only detectable 13CH
feature near 4211 Å to
estimate the isotope ratio 12C/13C.
The best fit model for the region
4210.7-4212.2 Å including also two 12CH
features was achieved for
12C/13C = 10
(Fig. 6,
middle panel).
However, because of the
of the spectrum of
HE 2327-5642 around 4211 Å, values of up to 12C/13C = 20
were found to be possible.

The abundance of nitrogen could only be determined from the NH
band at
3360 Å. In the literature, gf-values of
the NH molecular lines calculated
by Kurucz (1993) were subject
to corrections based on analysis of the Solar spectrum around
3360 Å. Hill et al.
(2002) apply a correction
of -0.807 in
to all of the NH lines, and Hayek
et al. (2009)
-0.4 dex. We checked the rather crowded spectral region around
3360 Å in the Solar spectrum
(Kurucz et al. 1984)
and fitted it with gf-values of the NH
lines that had been reduced by between -0.3
and -0.4 dex. With these corrections,
we derived a relative abundance, [N/Fe], of between -0.30
and -0.20
(Fig. 6,
bottom panel).

On the basis of its Li, C, and N abundances,
HE 2327-5642 does not appear to be exceptional.
Unfortunately, its oxygen abundance could not be determined from the
available
observed spectrum.

4.2 Sodium
to titanium

In HE 2327-5642, the -process elements Mg, Si, Ca,
and Ti are
enhanced relative to iron:
,
,
,
and .
This is
consistent with the behavior of other metal-poor halo stars (see, e.g.,
Cayrel et al. 2004).

The determination of the abundances of Mg and Ca is
based on NLTE line formation calculations for Mg I
and Ca I-II,
using the methods described by Zhao
et al. (1998, Mg I)
and
Mashonkina
et al. (2007a, Ca I-II).
For both elements, the same scaling factor,
,
was applied to the Steenbock &
Holweger (1984) formula for calculations of the inelastic
collisions with hydrogen atoms. Neutral Mg and Ca are minority species
in
the atmosphere of HE 2327-5642, and they are both subject to
overionization caused by super-thermal ultraviolet radiation of
non-local
origin, resulting in a weakening of the Mg I
and Ca I lines
relative to their LTE strengths. The NLTE abundance corrections,
,
are in the range 0.08-0.12 dex for the
Mg I lines and between 0.17 and
0.29 dex for the Ca I lines
(Table 8,
online material).

The Si abundance was derived from the only detected line, Si I 3905 Å,
assuming LTE. Based on the NLTE calculations for Si I
presented by Shi et al. (2009),
we estimated the NLTE abundance correction for
this line to be positive and on the order of a few hundredths of a dex.

Titanium is observed in HE 2327-5642 for two
ionization stages, and its
abundance can be reliably determined. We obtained a difference in
absolute LTE
abundances of -0.07 dex between Ti I
and Ti II. Assuming that
the NLTE effects for Ti II are as small, as
is the case for
Fe II, and that they are of the same order
for Ti I as they are
for Fe I, we found that
dex.

HE 2327-5642 displays an underabundance of the odd-Z
elements Na and Al
relative to iron of
and
.
This is not exceptional
for a metal-poor halo star.
Sodium and aluminium were observed in HE 2327-5642 only in the
resonance lines of
their neutral species. The abundance determination was based on NLTE
line
formation for Na I and Al I,
using the methods described by
Mashonkina et al. (1993)
and Baumüller & Gehren (1996).
For both species,
was
adopted. The NLTE abundances derived from the
Na I 5890/5896 Å lines are
-0.39/-0.28 dex lower than the corresponding LTE values. In
contrast, the NLTE abundance derived
from Al I 3961 Å is
0.52 dex higher than the LTE value. It is worth noting that
the calculated
of the Na lines agree
within 0.05/0.02 dex with those given by Andrievsky
et al. (2007) in their
Table 2 for
,
,
and
values close to those of
HE 2327-5642, while we found that
for
Al I is 0.25 dex lower than
indicated by Andrievsky
et al. (2008) in their Fig. 2 for
similar stellar parameters. For the relative abundances in
HE 2327-5642,
we obtained an Al/Na ratio close to Solar (
)
and
very low odd/even-Z ratios (
;
).

To determine the abundance of Sc, we employed four lines of
the majority
species Sc II. For each line, hyperfine
structure splitting was taken
into account, using the HFS data of McWilliam
et al. (1995). Neglecting the HFS effect
led to an overestimation of the Sc abundance of 0.08 dex for
Sc II 4246 Å, the
strongest line in the wavelength ranges covered by our
spectra. We obtained
,
which is about 0.2 dex lower
than the corresponding mean value for the cool halo giants studied by Cayrel et al. (2004). The
difference can be at least partly explained by HFS not having been
taken into account by Cayrel
et al. (2004). NLTE calculations for Sc II
in
the Sun were performed by Zhang
et al. (2008), with the result that the departures
from LTE are small with negative NLTE abundance corrections of -0.06 to
-0.03 dex.

4.3
Iron-group elements and Zn

We determined the abundance of six elements in this group. For two of
them, Mn
and Co, their energy levels are affected by considerable hyper-fine
splitting,
and HFS was explicitly taken into account in our spectrum
synthesis calculations where HFS data were available (see
Table 8
online for references).

For Mn, we measured 0.36 dex lower abundances from
the Mn I resonance
lines at Å
than from the Mn I subordinate line at
4041 Å in HE 2327-5642. A similar effect was found
for the Cr I lines:
two lines originating from the ground state, 4254 Å and
4274 Å, corresponded to 0.26 and 0.32 dex lower
abundances than the mean of the other chromium
lines. Our results are consistent with the findings of Johnson (2002) and later
studies. The Mn abundance derived from the Mn I
lines may be
underestimated because of departures from LTE. Bergemann
& Gehren (2008) predict
dex
for the Mn I resonance triplet in
the model 5000/4.0/-3, and 0.5 dex for Mn I
4041 Å. Usually,
the NLTE effects become more pronounced with decreasing .
However, it is unclear
whether
will vary with surface gravity in similar ways for the Mn I
resonance triplet and Mn I 4041 Å.
Therefore, the abundances derived from the Mn I
resonance lines were
not taken into account in calculating the mean presented in
Table 5.

We fortunately detected lines of Mn II,
the majority species of Mn,
which is hardly expected to be affected by departures from LTE,
according
to the results of Bergemann
& Gehren (2007).
We note that the relative LTE abundances [Mn I
(4041 Å)/Fe I] and
[Mn II/Fe II] in HE 2327-5642 are consistent
with each other to within
0.01 dex. Though HFS was not taken into account for the
Mn I 4041 Å line, its effect on
the abundance is expected to be
small, as the line is very weak (
mÅ). We measure for
HE 2327-5642 an underabundance of Cr and Mn very similar to
that of the comparison
sample (Cayrel et al. 2004).

HE 2327-5642 is also deficient in V and Ni relative
to iron and Solar
ratios. Information about V abundances in very metal-poor stars is
scarce in
the literature, probably due to difficulties in detecting the vanadium
lines.
We used four lines of V II located in the
blue spectral range, where
severe blending effects are present even in very metal-poor stars.
Paper II
found V/Fe ratios close to Solar for the sample covering a [Fe/H] range
from
-1.5 to -3. However, they noted that the V abundances are based on
quite
weak features and hence are susceptible to overestimation due to
unresolved
blends. For Ni, we used eight well observed and unblended lines of
Ni I. The large scatter in the abundances
obtained may be partly caused by our using four different sources for
the gf-values (see Table 8 online for
references). For example, the mean abundance derived
from two lines using the gf-values of Fuhr et al. (1988)
is 0.21 dex higher than the abundances measured from three
lines employing the gf-values of Blackwell et al.
(1989).

For the cobalt and zinc abundances of HE 2327-5642,
we obtained
values close to the Solar ratios with respect to iron, i.e.,
and
.
Using
from
the NLTE
calculations of Takeda et al.
(2005), and assuming similar departures from LTE for
Zn I 4722 Å, we calculated a NLTE
abundance ratio of
.
Cayrel et al. (2004)
found that [Co/Fe] and [Zn/Fe]
increase with decreasing metallicity, and measured
and
for
stars with [Fe/H] close to -2.8 (see their
Fig. 12). We note that Cayrel
et al. (2004) neglected HFS of the used lines
of Co I, thus they probably overestimated
the Co abundances.
According to our estimate for Co I
4121 Å in the atmospheric model
5050/2.34/-2.78, ignoring HFS makes a difference in abundance of
+0.09 dex.

Figure 7:

The heavy-element abundance pattern of HE 2327-5642 (filled
circles) compared to the Solar System r-process
(SSr) abundance pattern (continuous curve) scaled to match Ba-Hf. For
comparison, the heavy element abundances of the benchmark r-II stars
CS 22892-052 (open triangles), CS 31082-001
(crosses), and HE 1219-0312 (open circles) are shown. They
have been normalized to the value derived for
in
HE 2327-5642. The bottom panel displays the difference
between HE 2327-5642 and SSr defined as
.

4.4 Heavy elements

In the ``snapshot'' spectra of HE 2327-5642, Barklem et al.
(2005) detected
only six heavy elements beyond strontium. Owing to the higher quality
and
broader wavelength coverage of the spectra used in this study, we
detected 23 elements in the nuclear charge range between Z
= 38 and 90. We were unsuccessful in obtaining abundances for
Ru, Rh, and U. The Ru I 3436,
3728 Å and Rh I 3434,
3692 Å lines are very weak and
therefore could not be detected in our spectra. We marginally detected
the U II 3859.57 Å line
in our spectra of HE 2327-5642;
however, the S/N is not high
enough to derive a reliable abundance.

Our NLTE calculations for HE 2327-5642 showed that
the Sr II and
Ba II resonance lines are stronger than in
the LTE case,
resulting in negative NLTE abundance corrections of -0.15 dex
and
-0.09 dex, respectively. The subordinate lines of Ba II
exhibit a
different behavior: the weakest line at 5853 Å is weaker,
while the other two
lines, at 6141 and 6496 Å, are stronger relative to their
values for LTE.

In contrast to Sr II and Ba II,
the term structure of the other
NLTE species is produced by multiple electronic configurations and
consists of
hundreds and thousands of energy levels. For each of these species,
enhanced
photoexcitation from the ground state leads to overpopulation of the
excited
levels in the line formation layers, resulting in a weakening of the
lines. We
calculated positive NLTE abundance corrections for the lines of Zr II,
Pr II, and Eu II,
finding values close to +0.1 dex.
All the elements beyond barium are observed in the lines of their
majority
species, with term structures as complicated as that for Eu II,
so
the departures from LTE are expected to be similar to those for Eu II.
This is largely true also for osmium and iridium detected in the lines
of
their neutrals, Os I and Ir I,
which have relatively high
ionization energies of 8.44 and 8.97 eV, respectively.
Fortunately, the
abundance ratios among heavy elements are
probably only weakly affected
by departures from LTE. For consistency, we used in this study the
abundances
of the heavy elements beyond strontium as determined based on the LTE
assumption.
They are presented in Table 5 and
Fig. 7.

We now explore the abundance patterns of elements in the three
r-process
peaks.

Five elements with
were
measured in the region of the first
peak. The only molybdenum line in the visible spectrum, Mo I 3864 Å,
can be used to determine the element abundance of cool stars. In
HE 2327-5642, this line is nearly free of blends, but it is
weak: the
central depth of the line is only 2% of the continuum. With
of
the observed spectrum within this wavelength region, the uncertainty
in the derived Mo abundance was estimated to be 0.3 dex.

A similar uncertainty is expected for palladium, which was
detected in a
single line, Pd I 3404 Å.
In HE 2327-5642, this line is free of
blends and is stronger than that of Mo I 3864 Å,
but it is located
in a spectral range where the S/N
is only 20 to 25.

Strontium is observed in HE 2327-5642 in two strong
resonance lines,
Sr II 4077
and 4215 Å. Both lines are affected by HFS of the odd
isotope .
The synthetic spectrum was calculated with the
fraction
of 0.22 corresponding to a pure r-process
production of strontium (Arlandini
et al. 1999, stellar model).

The best fits (continuous curve) of Ho II
3456 Å ( top panel) and Hf II
3399 Å ( bottom panel) in the observed
spectrum of HE 2327-5642 (bold dots). The dashed curves show
the effect of a 0.1 dex variation in the abundance on the
synthetic spectrum. The dotted curves show the synthetic spectrum with
no holmium and hafnium in the atmosphere.

With 15 elements measured in the Ba-Hf range (see Fig. 8
for holmium and hafnium), the second r-process peak
is the most tightly constrained among the three peaks.

The barium abundance given in Table 5 was
determined
from the three subordinate lines, Ba II
5853, 6141, and 6497 Å,
which are almost free of HFS effects. According to our estimate for
Ba II 6497 Å, neglecting HFS makes
a difference in abundance of no
more than 0.01 dex. In contrast, the Ba II 4554 Å
resonance line
is strongly affected by HFS. The even isotopes are unaffected by HFS,
while the odd isotopes exhibit significant HFS, and thus the element
abundance derived from this line depends on the Ba isotope mixture
adopted in the calculations. Since the odd isotopes
and
have
very similar HFS, the abundance is essentially dependent on the
total fractional abundance of these odd isotopes,
.
For example, using the Solar Ba isotope mixture with
,
we inferred an abundance that was 0.34 dex higher from Ba II 4554 Å
than the mean abundance of the subordinate lines. In the LTE
calculations, the difference was reduced by 0.20 dex when we
adopted a pure r-process Ba isotope mixture with
,
as predicted by
Arlandini et al. (1999, stellar
model).
The remaining discrepancy between the resonance
and subordinate lines was largely removed by NLTE calculations. Thus,
our
analysis of HFS affecting the Ba II 4554 Å
resonance line suggests
a pure r-process origin of barium in
HE 2327-5642.

The lines of Eu II and Yb II
detected in HE 2327-5642
consist of multiple IS and HFS components. To derive the total
abundance of the
given element, we adopted in our calculations a pure r-process
isotope mixture
from the predictions of Arlandini
et al. (1999, stellar model):
:
= 39:61 and
:
:
:
:
=
18.3 : 22.7 : 18.9 : 23.8 : 16.3. All the
lines of Nd II and Sm II
observed in HE 2327-5642 are rather
weak, and are treated as single lines.

Only one hafnium line can be measured in
HE 2327-5642. The
observed feature is attributed to a combination of the Hf II
3399.79 Å and NH 3399.79 Å molecular line. Using
taken
from Kurucz (1993) and reduced
by
-0.3 dex, and the nitrogen abundance
,
the molecular line contributes
approximately 45 % to the 3399 Å blend (Fig. 8).
Ignoring the molecular contaminant completely leads to a
0.17 dex higher
hafnium abundance. We therefore estimate the uncertainty in the Hf
abundance obtained to be 0.15 dex.

The third peak and actinides were probed for three elements,
osmium, iridium,
and thorium. The abundance of osmium was determined from a single line,
Os I 4260 Å. The line is
weak, with a center line depth of 2.5% in the continuum flux, and is
nearly free of blends. Because of the high
of
the observed spectrum, the uncertainty in the
derived osmium abundance was estimated as 0.2 dex.

Two iridium lines, Ir I 3513 and
3800 Å, were clearly detected in HE 2327-5642
(Fig. 9).
The theoretical profiles were calculated by
taking HFS effects into account and iridium isotope abundance ratio
:
= 37 : 63,
which is obtained to be the same in the
matter for the Solar System (Lodders
2003) and the matter produced by the
r-process (Arlandini
et al. 1999). The Ir I
3800 Å line was measured in two observed spectra, BLUE390 and
437BLUE, (Table 2)
and seems reasonably to be reliable. With its equivalent width of
5.8 mÅ and
of the observed spectrum, the uncertainty in the
derived iridium abundance was estimated to be 0.08 dex. The Ir
I 3513 Å line served as a
verification that it agrees with the other line. The blend at
3513.6 Å is well reproduced by assuming the value of
found
from Ir I 3800 Å, as shown
in Fig. 9.

The radioactive element thorium was clearly detected in
HE 2327-5642 in the Th II
4019 Å line (Fig. 10), but
proved rather challenging to incorporate into the
determination of the stellar age of HE 2327-5642. Unhappily,
the observed spectrum around Th II
4019 Å has low signal-to-noise ratio (
), and the uncertainty in the
derived thorium abundance was estimated to be 0.2 dex.

4.5 Error
budget

We performed a detailed error analysis of HE 2327-5642, to
estimate the uncertainties in the abundance measurements for the heavy
elements beyond the iron group. Stochastic errors (
)
caused by random uncertainties in the continuum placement, line profile
fitting,
and gf-values, are represented by a dispersion in
the measurements of
multiple lines around the mean (
),
as given in
Table 5
when
lines of an element are
observed. Observational errors in the species with a single line used
in
abundance analysis were discussed in Sect. 4.4. Systematic
uncertainties include those that exist in the adopted stellar
parameters, in
the used hydrostatic model atmospheres, and in the LTE line formation
calculations. As argued in Sect. 4.4, the latter
is not expected
to influence the abundance pattern of the elements in the range from La
to Th.

It is difficult to estimate the uncertainty introduced by
using the 1D model
atmosphere. The elements in the La-Th range, for example, are observed
in the lines of their majority species. The detected lines arise from
either the
ground or low-excitation levels, and most of them are relatively weak,
i.e.,
mÅ. This means that
they all are formed in the same
atmospheric layers. It would therefore be rather unexpected for there
to be significantly different 3D
effects for individual elements in the La-Th range. This is different
from the case of the elements in the Sr-Pd range, Ba, and Yb, which are
observed
in either the strong lines (e.g. Sr II 4215 Å,
Ba II 6141 Å,
Yb II 3694 Å) or the
lines of the minority species Mo I and
Pd I. Both NLTE and 3D effects may have a
strong influence on their
derived abundance. Hence, we examined here only those uncertainties
linked to
our choice of stellar parameters. These were estimated by varying
by
-70 K,
by -0.1 dex, and
by -0.1 km s-1 in the stellar
atmosphere model.

Table 6
summarizes the various sources of
uncertainties. The quantity
listed
in Col. 5 is the total
impact of varying each of the three parameters, computed as the
quadratic sum
of Cols. 2-4. The total uncertainty
(Col. 7)
in the
absolute abundance of each element is computed by the quadratic sum of
the
stochastic (
)
and systematic (
)
errors.

5 Comparison to other r-II stars and Solar r-process
abundances

The abundance pattern of the neutron-capture elements in the range from
Sr to
Os in HE 2327-5642 is very similar to that of other
well-studied r-II
stars. Figure 7
shows comparisons with CS 22892-052
(Sneden et al. 2003),
CS 31082-001 (Hill
et al. 2002; Plez et al. 2004), and
HE 1219-0312
(Hayek et al. 2009).
For example, the dispersion about the mean of the quantities
amounts
to 0.10 dex, which is
at the level of
error bars in the abundance determinations of the
individual elements.

Figure 11:

The heavy-element abundance pattern of HE 2327-5642 (filled
circles) compared to the r-residuals calculated
with various Solar total abundances and s-process
abundances. Continuous and dotted curves correspond to the predicted s-process
abundances of Arlandini et al.
(1999, stellar model) and the Solar total abundances from Lodders et al. (2009)
and Asplund et al. (2009),
respectively. The dashed curve (only Sr-Mo) corresponds to the s-process
abundances of Travaglio
et al. (2004) and the Solar total abundances of Lodders et al. (2009).
Each r-process abundance pattern was scaled to
match the Ba-Hf in HE 2327-5642.

Iridium in HE 2327-5642 seems to be overabundant with respect
to
CS 22892-052 and CS 31082-001. Roederer et al. (2009)
compiled and/or revised the
neutron-capture element abundances for a sample of r-process
rich stars
including CS 22892-052 and CS 31082-001, and they
obtained a mean iridium to
europium ratio of
for 15
stars.
For HE 2327-5642, we derived
.
We
investigated the possible sources of the difference between our
analysis and that
of Roederer et al. (2009),
identifying one that could at least partly explain it.
The iridium abundances of Roederer
et al. (2009) were underestimated
by approximately 0.2 dex because they used the partition
functions of
Ir II implemented in the MOOG code (C.
Sneden, private
communication), which are about a factor of two lower than that we
used in our study. We compared the partition functions thoroughly and
are satisfied that our new data, which are based on the highest quality
energy level data available, are more accurate. In all other relevant
respects (i.e., gf-values, HFS
data, ionization potential of Ir I), our
data and methods and those of
Roederer et al. (2009)
are identical. The difference in the codes used for the
abundance determination cannot play a role only for iridium, leaving
all the
other determinations unaffected. Hence, we are left with a discrepancy
of
0.14 dex in
between
HE 2327-5642 and the
stellar sample of Roederer
et al. (2009). However, to draw any firm
conclusion at this point, the abundance determinations have to be
confirmed using more accurate measurements and the latest detections of
the third r-process peak
elements.

The elements in the range from Ba to Hf in
HE 2327-5642 were found to match
the scaled Solar r-process pattern very well, with
a dispersion of 0.07 dex
about the mean of the differences
.
This
is in line with earlier results obtained for other r-process
rich stars,
e.g., CS 22892-052 (Sneden
et al. 2003), CS 31082-001 (Hill et al. 2002),
HD 221170
(Ivans et al. 2006),
CS 29491-069, and HE 1219-0312 (Hayek
et al. 2009), and provides additional evidence of
universal production ratio of these elements during
the Galactic history. It is worth noting that the use of the partition
function of Ho II from Bord
& Cowley (2002) improves the comparison with the
scaled Solar r-process for holmium.

For the lighter elements in HE 2327-5642, the
difference
is
indicative of a large spread of the data, between
-0.4 dex (Mo) and +0.7 dex (Y). In the following, we
show that at
least a major fraction of the departures from the Solar r-process
found for
the light trans-iron elements is likely to be due to inaccurate Solar
r-residuals.

For a given element, the r-residual is
obtained by subtracting theoretical
s-process yields from the observed total Solar
abundance. We consider, for instance,
the s-process abundances from Arlandini
et al. (1999, stellar model) and use the
Solar total abundances from two different sources,
Lodders et al. (2009,
meteoritic) and Asplund
et al. (2009, photospheric). The
0.02 dex increase in the yttrium abundance as one changes from
Lodders et al. (2009)
to the Asplund et al. (2009)
data leads to a 0.65 dex increase
in the r-residual. A notable difference between two
sets of the Solar
r-process abundances was found for all elements with
dominant s-process
contribution to their Solar abundances, for example, -0.58 dex
for Sr,
-0.40 dex for Zr, and -0.31 dex for La (see
Fig. 11).
This
is because the calculation of the r-residuals
involves the subtraction of a
large number from another large number, so that any small variations of
one of
them leads to a dramatic change in the difference.

Significant uncertainties in the r-residual
are also caused by differences
between s-process calculations. For example, Arlandini et al. (1999)
obtained the s-process
abundance distribution by performing a best-fit to the Solar mains-component
using stellar AGB models of 1.5
and 3
with half-Solar
metallicity. Travaglio
et al. (2004,1999) calculated the s-process
contribution to the Solar abundances by integrating of the s-process
yields
of different generations of AGB stars, i.e., considering the whole
range of
Galactic metallicities. In both studies, very similar results were
found for
Ba and Eu; however, Travaglio
et al. (2004) predicted lower s-process
abundances
for the elements in the Sr-Mo range. In consequence, their
Solar r-residuals
were significantly increased, as shown in Fig. 11. From this
discussion, it is clear that no solid conclusion can be drawn about any
departures from the scaled Solar r-process pattern
in the
Sr-Pd region for r-II stars.

As expected, all stars within each group have very similar
Ba/Eu abundance ratios,
as shown in the bottom panel of Fig. 12, where mean
Ba/Eu ratios of
(r-II stars),
(r-I
stars),
and
(Eu-poor stars). We note that
for
a pure r-process production of heavy elements (Arlandini et al. 1999).
This suggests that only a
small number of s-nuclei (including those of
strontium) existed in the
matter out of which these stars formed. The Sr/Eu abundance ratios
indicate a
clear separation between each of these three groups (Fig. 12).
We note that the Zr/Eu ratios exhibit very similar behavior. The mean
Zr/Eu
abundance ratios are
(r-II
stars),
(r-I stars), and
(Eu-poor stars).
HE 2327-5642, having
,
(crossed
circle in
Fig. 12),
and ,
clearly
belongs to the group of r-II stars.

From Figs. 7
and 12,
it is clear that the
first and second r-process peak elements in the
r-II stars are of common
origin. However, the origin of the first neutron-capture peak elements
in the
r-I and Eu-poor stars remains unclear, despite being the subject of a
number of studies
(Travaglio
et al. 2004; Farouqi et al. 2009; Truran
et al. 2002).

Figure 12:

The Sr/Eu (top panel) and Ba/Eu ( bottom
panel) abundance ratios in r-II (filled circles), r-I (open
rombs), and Eu-poor (asterisks) stars (for the sources of the data, see
text). HE 2327-5642 is shown by a crossed open circle. The
solid and dotted lines indicate the pure r-process
and Solar system ratios, respectively.

6 Age determination

The detection of thorium permits a nucleo-chronometric age
estimation of HE 2327-5642 by comparing the observed
Th-to-stable neutron-capture element-abundance ratios with the
corresponding initial values at the time when the star was born, (Th/X)0:

Because of the uncertainty in the thorium abundance of
0.2 dex, which
translates into an age uncertainty of 9.3 Gyr, a precise age
estimation is
not possible. Nevertheless, we investigated how the results depend on
the
adopted values of (Th/X)0,
and which Th/X pairs are possibly
reliable chronometers. We assumed that thorium was produced with the
lighter elements in the range between Ba and Hf.

We first determined the age of the star using the initial
abundance ratios
from the dynamical network calculations of Farouqi
et al. (2008). They considered a
core-collapse supernova (SN II) with an adiabatically
expanding high-entropy
wind (HEW) as the astrophysical environment for the r-process.
In the HEW
scenario, the total nucleosynthetic yield is the sum of SN ejecta with
multiple components in different entropy ranges. Farouqi
et al. (2008) found that
heavy elements beyond
are produced in the highest entropy (S >
150) zones in the so-called ``main'' r-process. The
HEW model production
ratios, ,
as given by Hayek et al.
(2009),
and the calculated ages for multiple Th/X pairs
(where X is one of the
elements in the Ba-Ir range for which we determined an abundance) are
listed
in Table 7.

The age uncertainties introduced by the measurement
uncertainties are listed
in the column ``Error''. They were calculated as
Gyr,
where
is taken from Table 6.
Variations in the stellar parameters
and
yielded an
uncertainty of 1.5 Gyr in the final age. It
can be seen that individual pairs indicate a large spread in stellar
age. The
mean value is Gyr.
As noted by Hayek et al.
(2009), the
estimate for hafnium in the HEW model is rather uncertain due to
problems with the
nuclear data. Neglecting hafnium and also osmium, because of less
reliable stellar abundance, yields
Gyr,
which agrees well with the expected age of an extremely metal-poor star
that formed in the early Galaxy. We note that the cosmic age derived
from the results of the Wilkinson Microwave Anisotropy Probe
(WMAP) experiment is
Gyr
(Spergel et al. 2003).

For an additional estimate of the age of
HE 2327-5642, we employed the
Solar r-residual ratios
for
the elements for which the
r-process fraction exceeds 70% (columns SSr
in
Table 7).
Since these are measured values, they
depend only weakly on theoretical predictions and their associated
nuclear
physics uncertainties. Since the Sun is approximately 4.5 Gyr
old, the
corresponding correction accounting for the thorium radioactive decay
was
introduced to the Solar current thorium abundance. The resulting mean
age of
HE 2327-5642, Gyr,
calculated using all stable elements from Sm to Ir seems low for a halo
star. If we neglect the estimate based on Th/Ir, which is clearly an
outlier (i.e.,
Gyr), we obtain an
even lower stellar age of
Gyr.

We note that the stochastic error in the stellar age based on
the HEW model predictions is large compared to that for the Solar r-residual
ratios. This is probably caused by the
uncertainty in the theoretical yields for individual elements.

Table 7:
Logarithmic production ratios (PR)
for the HEW model and Solar System r-process (SSr)
and corresponding radioactive decay ages in HE 2327-5642.

7 Conclusions

The high-quality VLT/UVES spectra of HE 2327-5642 has enabled
us to
determine accurate abundances for 40 elements, including 23 elements
in the nuclear charge range Z
= 38-90. We have confirmed that
HE 2327-5642 is strongly r-process
enhanced, having
where
r denotes the average of the abundances of
seven
elements (i.e., Eu, Gd, Tb, Dy, Ho, Er, and Tm), where there is an r-process
contribution to
the Solar system matter of more than 85 % according to the r-residuals
of
Arlandini et al. (1999).
We have found that HE 2327-5642 and three benchmark r-II
stars,
i.e., CS 22892-052 (Sneden
et al. 2003), CS 31082-001 (Hill et al. 2002), and
HE 1219-0312
(Hayek et al. 2009)
have very similar abundance patterns of the elements in the range
from Sr to Os. Hence, HE 2327-5642 is a member of the small
sample of
currently known r-II stars.

The elements in the range from Ba to Hf in
HE 2327-5642 match the scaled
Solar r-process pattern very well. We have shown
that the Solar r-residuals
for the first r-process peak elements are rather
uncertain. They may vary by as
much as 0.5 dex or even more, depending on the adopted Solar
total abundances
and s-process fractions. Therefore, no firm
conclusion can be drawn about any relation between the light trans-iron
elements in r-II stars and the Solar r-process.

We have found a clear distinction in Sr/Eu abundance ratios
between the halo stars
of different europium enhancement. We have proposed using the [Sr/Eu]
ratio in
addition to [Eu/Fe] to separate the strongly r-process
enhanced (r-II) stars
from the other halo stars that have experienced a dominant contribution
of the r-process to
heavy element production. The r-II stars, whose stellar matter
presumably has
experienced a single nucleosynthesis event, have
,
,
and a low Sr/Eu abundance ratio of
.
Stars with very similar Ba/Eu ratios have
two times (0.36 dex) higher Sr/Eu ratios if their Eu/Fe ratio
is in the range
(i.e.,
r-I stars), and nearly an order of
magnitude (0.93 dex) higher Sr/Eu ratios if
(Eu-poor
stars). The origin of the first neutron-capture peak elements in the
r-I stars
and Eu-poor stars remains unclear, and additional theoretical studies
are needed
to elucidate this problem.

Only two elements, Os and Ir, of the third r-process
peak were detected in
HE 2327-5642. Iridium appears to be overabundant compared to
the Ir
abundance determined in other r-process enhanced
stars. However, due to the
uncertainty in the Ir abundance, we have been unable to draw a firm
conclusion.

The detection of thorium permitted an estimate of the
radioactive decay age of
HE 2327-5642, although the age uncertainty of 9.3 Gyr
introduced by the uncertainty in the thorium abundance is rather large.
Employing multiple Th/X chronometers and initial
production ratios based on the Solar r-residuals,
an age of Gyr
was obtained from nine Th/X pairs, involving
elements in the
Sm-Os range. Using the predictions of the HEW r-process
model, as given by
Hayek et al. (2009),
we obtained Gyr
from 12 Th/X pairs.

Based on our high-resolution spectra, covering 4.3 years,
we propose that HE 2327-5642 is a radial-velocity variable
with a highly elliptical orbit of the system. Determination of the
orbital period would provide the unique opportunity to determine a
lower limit to the mass of the secondary in this system. Scenarios for
the site of the r-process include a high-entropy
wind from a type-II supernova (e.g., Woosley et al. 1994; Takahashi
et al. 1994), ejecta from neutron star mergers (e.g.
Freiburghaus et al. 1999),
or the neutrino-driven wind of a newly formed neutron star in an
accretion-induced collapse (AIC) event (e.g. Woosley & Baron 1992; Qian &
Wasserburg 2003). According to these scenarios, it is
expected that the secondary is a neutron-star. With a lower limit to
the mass of the secondary, it might be possible to constrain a
scenario, because in the AIC case the neutron star is expected to have
a mass just slightly above the Chandrasekhar mass, while core-collapse
supernovae or neutron star mergers would result in remnants of
significantly higher mass.

Acknowledgements

The authors thank Thomas Gehren for the NLTE calculations
for Al I and Tatyana Ryabchikova for help
with collecting the atomic data. L.M. and A.V. are supported by the
Russian Foundation for Basic Research (grant 08-02-92203-GFEN-a), the
Russian Federal Agency on Science and Innovation
(No. 02.740.11.0247), and the Swiss National Science
Foundation (SCOPES project No. IZ73Z0-128180/1). N.C. is
supported by the Knut and Alice Wallenberg Foundation, and by Deutsche
Forschungsgemeinschaft through grants Ch 214/3 and
Re 353/44.
P.S.B. is a Royal Swedish Academy of Sciences Research Fellow supported
by a grant from the Knut and Alice Wallenberg Foundation. P.S.B. also
acknowledges additional support from the Swedish Research Council.
T.C.B. acknowledges partial
funding of this work from grants PHY 02-16783 and PHY 08- 22648:
Physics Frontier Center/Joint Institute for Nuclear Astrophysics
(JINA), awarded by the U.S. National Science Foundation.
We made use of model atmosphere from the MARCS library, and the NIST
and VALD databases.

Online Material

Table 8:
Line data and abundances from an analysis of HE 2327-5642.
corresponds to 10 000 K. Column 6 gives references to
the adopted gf-values. Column 13 gives references
to the sources of the used IS and HFS data and adopted values.

Footnotes

... HE 2327-5642

Based on
observations collected at the European Southern Observatory, Paranal,
Chile (Proposal numbers 170.D-0010, and 280.D-5011).

Table 7:
Logarithmic production ratios (PR)
for the HEW model and Solar System r-process (SSr)
and corresponding radioactive decay ages in HE 2327-5642.

Table 8:
Line data and abundances from an analysis of HE 2327-5642.
corresponds to 10 000 K. Column 6 gives references to
the adopted gf-values. Column 13 gives references
to the sources of the used IS and HFS data and adopted values.

All Figures

Synthetic profiles of H
( top panel) and H
( bottom panel) from the s-s
(dashed curve), s-p (continuous
curve), and p-p (dotted curve)
models. The calculations for the s-p
and p-p models were made for
pure hydrogen lines.

Top panel: synthetic flux profile of H computed
for K
(continuous curve) compared to the observed spectrum of
HE 2327-5642 (bold dots). The dashed curves show the effect of
a 80 K variation in the effective temperature on the synthetic
spectrum. In all calculations, we assumed
,
,
and km s-1.
Bottom panel: effective temperature derived from the H
(filled circles) and H
(open diamonds) line wings in HE 2327-5642 as a function of
surface gravity. The error bars show the uncertainty of
arising
from profile fitting.

NLTE abundances of Fe I (filled circles),
Fe II (open circles), Ca I
(filled diamonds), and Ca II (open
diamonds) in HE 2327-5642 as a function of surface gravity.
For clearer illustration, the symbols for Ca are shifted upwards by
0.5 dex. The calculations are for
K,
,
and km s-1.

Trends of abundances with equivalent width and excitation potential, as
determined from individual Fe I (filled
circles) and Fe II (open circles) lines,
using our adopted stellar parameters. The dotted line indicates the
mean Fe abundance from two ionization stages and the shaded grey area
its statistical error.

Best fits (continuous curve) of the CH features near 4310 Å (
top panel) and 4211 Å ( middle panel),
and the NH molecular band near 3360 Å ( bottom panel).
The observed spectrum of HE 2327-5642 is shown as bold dots.
The dashed curves in the top and bottom panels
show the synthetic spectra with no carbon and nitrogen in the
atmosphere. In the middle panel, the continuous
curve corresponds to an isotope ratio of
,
while the dashed curves are synthetic spectra for
and 30.

The heavy-element abundance pattern of HE 2327-5642 (filled
circles) compared to the Solar System r-process
(SSr) abundance pattern (continuous curve) scaled to match Ba-Hf. For
comparison, the heavy element abundances of the benchmark r-II stars
CS 22892-052 (open triangles), CS 31082-001
(crosses), and HE 1219-0312 (open circles) are shown. They
have been normalized to the value derived for
in
HE 2327-5642. The bottom panel displays the difference
between HE 2327-5642 and SSr defined as
.

The best fits (continuous curve) of Ho II
3456 Å ( top panel) and Hf II
3399 Å ( bottom panel) in the observed
spectrum of HE 2327-5642 (bold dots). The dashed curves show
the effect of a 0.1 dex variation in the abundance on the
synthetic spectrum. The dotted curves show the synthetic spectrum with
no holmium and hafnium in the atmosphere.

The heavy-element abundance pattern of HE 2327-5642 (filled
circles) compared to the r-residuals calculated
with various Solar total abundances and s-process
abundances. Continuous and dotted curves correspond to the predicted s-process
abundances of Arlandini et al.
(1999, stellar model) and the Solar total abundances from Lodders et al. (2009)
and Asplund et al. (2009),
respectively. The dashed curve (only Sr-Mo) corresponds to the s-process
abundances of Travaglio
et al. (2004) and the Solar total abundances of Lodders et al. (2009).
Each r-process abundance pattern was scaled to
match the Ba-Hf in HE 2327-5642.

The Sr/Eu (top panel) and Ba/Eu ( bottom
panel) abundance ratios in r-II (filled circles), r-I (open
rombs), and Eu-poor (asterisks) stars (for the sources of the data, see
text). HE 2327-5642 is shown by a crossed open circle. The
solid and dotted lines indicate the pure r-process
and Solar system ratios, respectively.

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